Gaseous Planets, Protostars, and Young Brown Dwarfs: Birth and Fate
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Chabrier: Formation of Gaseous Planets, Protostars, and Young Brown Dwarfs 623 Gaseous Planets, Protostars, and Young Brown Dwarfs: Birth and Fate G. Chabrier, I. Baraffe, and F. Selsis Ecole Normale Supérieure de Lyon T. S. Barman University of California, Los Angeles P. Hennebelle Ecole Normale Supérieure, Paris Y. Alibert University of Bern We review recent theoretical progress aimed at understanding the formation and the early stages of evolution of giant planets, low-mass stars, and brown dwarfs. Calculations coupling giant planet formation, within a modern version of the core accretion model that includes planet migration and disk evolution, and subsequent evolution yield consistent determinations of the planet structure and evolution. Uncertainties in the initial conditions, however, translate into large uncertainties in the luminosity at early stages. It is thus not possible to say whether young planets are faint or bright compared with low-mass young brown dwarfs. We review the effects of irradiation and evaporation on the evolution of short-period planets and argue that substan- tial mass loss may have occurred for these objects. Concerning star formation, geometrical effects in protostar core collapse are examined by comparing one-dimensional and three-dimensional calculations. Spherical collapse is shown to significantly overestimate the core inner density and temperature and thus to yield incorrect initial conditions for pre-main-sequence or young brown dwarf evolution. Accretion is also shown to occur nonspherically over a very limited fraction of the protostar surface. Accretion affects the evolution of young brown dwarfs and yields more compact structures for a given mass and age, thus fainter luminosities, confirming previous studies for pre-main-sequence stars. This can lead to severe misinterpretations of the mass and/or age of young accreting objects from their location in the Hertzsprung-Russell (HR) diagram. Since accretion covers only a limited fraction of the protostar surface, we argue that newborn stars and brown dwarfs should appear rapidly over an extended area in the HR diagram, depending on their accretion history, rather than on a well-defined birth line. Finally, we suggest that the distinction between planets and brown dwarfs be based on an observational diagnostic, reflecting the different formation mechanisms between these two distinct populations, rather than on an arbitrary, confusing definition. 1. INTRODUCTION these objects are located well within the so-called ice line and could not have formed in situ. This strongly favors One of the fundamental questions of astrophysics re- planet migration as a common process in planet formation. mains the characterization of the formation of planets and This issue is explored in section 2 where we present con- stars. The mass ranges of the most massive planets and of sistent calculations between a revised version of the core the least-massive brown dwarfs certainly overlap in the ~1– accretion model, which does take planet migration into 10 MJup range; it is thus interesting to explore our under- account, and subsequent evolution. In this section, we also standing of the planet- and star-formation mechanisms in review our current understanding of the effects of irradia- a common review. tion and evaporation on the evolution of short-period plan- The growing number of discovered extrasolar giant plan- ets, hot Neptunes, and hot Jupiters, and review present un- ets, ranging now from Neptune-mass to few Jupiter-mass certainties in the determination of the evaporation rates. In objects, has questioned our understanding of planet forma- section 3, we briefly review our current understanding of tion and evolution. The significant fraction of exoplanets in protostellar core collapse and we show that nonspherical close orbit to their parent star, in particular, implies a revi- calculations are required to obtain proper accretion histories, sion of our standard scenario of planet formation. Indeed, densities, and thermal profiles for the prestellar core. The 623 624 Protostars and Planets V effect of accretion on the early contracting phase of pre- phases 1 and 2, and phase 2 essentially determines the for- main-sequence stars and young brown dwarfs, and a review mation timescale of the planet. The planet can thus form of observational determinations of accretion rates, are con- now on a timescale consistent with disk lifetimes, i.e., a few sidered in section 4. Finally, throughout this review, we have million years for a Jupiter (see A05). adopted as the definition of planet an object formed by the In the models of Bodenheimer et al. (2000a) and Hu- three-step process described in section 2.1, characterized by byckij et al. (2005), which are based on the P96 formal- a central rocky/icy core built by accretion of planetesimals ism, the calculations proceed in three steps: (1) the planet in a protostellar nebula. In contrast, genuine brown dwarfs is bounded by its Roche lobe (Rp = RL) [or more precisely 2 are defined in this review as gaseous objects of similar com- by Min(RL,Racc) where Racc = GM/cs is the accretion radius position as the parent cloud from which they formed by col- and cs the local sound velocity in the disk] so that the tem- lapse. This issue is discussed in section 5 and observational perature and pressure at the planet surface are the ones of diagnostics to differentiate brown dwarfs from planets, based the surrounding nebula. Note that in P96 calculations, opac- on their different formation mechanisms, are suggested. Sec- ity of the nebula is a key ingredient. (2) The planet external tion 6 is devoted to the conclusion. radius is the one obtained when the maximum gas accretion rate is reached. In P96, this value is fixed to 1 × 10 –2 M yr –1. 2. GASEOUS PLANETS: BIRTH At this stage, the external conditions have changed (Rp < AND EVOLUTION RL). Matter falls in free fall from the Roche lobe to the planet radius, producing a shock luminosity. (3) Once the 2.1. Planet Formation planet reaches its predefined final mass, the accretion rate is set to 0 and the boundary conditions become the ones of πσ 2 4 κ 2 The conventional planet formation model is the core ac- a cooling isolated object, L = 4 R Teff and RPph = 3g, κ cretion model as developed by Pollack et al. (1996, here- where R denotes the mean Rosseland opacity. The planet after P96). One of the major difficulties faced by this model surface radius is essentially fixed by the accretion shock is the long timescale necessary to form a gaseous planet like conditions (see, e.g., Fig. 1d of Hubickyj et al., 2005). This Jupiter, a timescale significantly larger than typical disk life- value, however, remains highly uncertain, as its correct de- times, <10 m.y. Reasonable timescales can be achieved only termination would imply a proper treatment of the radiative at the expense of arbitrary assumptions such as, e.g., nebula shock. In A05, phase 1 is similar to step 1 described above, mean opacities reduced to 2% of the ISM value in some tem- except that the planet migration from an initial arbitrary perature range or solid surface density significantly larger location and the disk evolution are taken into account, so than the minimum mass solar nebula value (Hubickyj et al., that the thermodynamic conditions of the surrounding neb- 2005). This leaves the standard core accretion model in an ula, as well as the distance to the star and thus the planet uncomfortable situation. This model has been extended re- Roche lobe radius, change with time. The planet’s final mass cently by Alibert et al. (2004, 2005, hereafter A05) by in- is set by the accretion rate limit, and is thus not defined a cluding the effects of migration and disk evolution during the priori. Note that, because of the disk evolution and/or the planet-formation process. The occurence of migration dur- creation of a gap around the planet, the accretion rate limit ing planet formation is supported by the discovery of nu- is 1 to 2 orders of magnitude smaller than the one in P96 merous extrasolar giant planets at very short distance to their at the end of phase 1 and reaches essentially 0 with time, a parent stars, well within the so-called ice line, about 5 AU fact supported by three-dimensional hydrodynamical simu- for the solar nebula conditions. Below this limit, above ice lations (D’Angelo et al., 2003; Kley and Dirksen, 2006). melting temperature, the insufficient surface density of sol- Eventually the planet opens a gap when its Hill radius be- ids that will eventually form the planet core, and the lack comes equal to the disk density scale height and migration of a large reservoir of gas, prevent in situ formation of large stops or declines until the disk is dissipated (see A05 for gaseous planets. details). The planet radius cannot be defined precisely in this Moreover, inward migration of the planet should arise model as it results from the competing effects of gas accre- from angular momentum transfer due to gravitational in- tion and planet contraction with changing boundary condi- teractions between the gaseous disk and the growing planet tions as the planet migrates inward and the disk evolves. In (Lin and Papaloizou, 1986; Ward, 1997; Tanaka et al., any event, the final stages of accretion are likely to occur 2002). Taking into account the migration of a growing within streams (see, e.g., Lubow et al., 1999), i.e., non- planet solves the long-lasting timescale problem of the core- spherically and, as mentioned above, the planet final radius accretion scenario. Indeed, when migration is included, the remains highly uncertain, at least in any one-dimensional planet-feeding zone never becomes depleted in planetesi- calculation.